2020-04-29T16:05:46Z
Acceptance in OA@INAF
High-resolution abundance analysis of red giants in the metal-poor bulge globular
cluster HP 1
Title
Barbuy, B.; Cantelli, E.; Vemado, A.; Ernandes, H.; Ortolani, S.; et al.
Authors
10.1051/0004-6361/201628106
DOI
http://hdl.handle.net/20.500.12386/24344
Handle
ASTRONOMY & ASTROPHYSICS
Journal
591
Number
DOI:10.1051/0004-6361/201628106
c ESO 2016
Astronomy
Astrophysics
&
High-resolution abundance analysis of red giants in the metal-poor
bulge globular cluster HP 1
?
B. Barbuy
1, E. Cantelli
1, A. Vemado
1, H. Ernandes
1, S. Ortolani
2, 3, I. Saviane
4, E. Bica
5, D. Minniti
6, 7, B. Dias
4,
Y. Momany
3, V. Hill
8, M. Zoccali
6, 9, and C. Siqueira-Mello
11 Universidade de São Paulo, IAG, Rua do Matão 1226, Cidade Universitária, 05508-900 São Paulo, Brazil
e-mail: barbuy@astro.iag.usp.br
2 Dipartimento di Fisica e Astronomia, Università di Padova, 35122 Padova, Italy
3 INAF–Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122 Padova, Italy 4 European Southern Observatory, Alonso de Cordova 3107, Santiago, Chile
5 Universidade Federal do Rio Grande do Sul, Departamento de Astronomia, CP 15051, 91501-970 Porto Alegre, Brazil 6 Millenium Institute of Astrophysics, Av. Vicuña Mackenna 4860, Macul, Santiago, Chile
7 Departamento de Ciencias Fisicas, Universidad Andres Bello, Republica 220, Santiago, Chile
8 Laboratoire Lagrange (UMR 7293), Université de Nice Sophia Antipolis, CNRS, Observatoire de la Côte d’Azur, CS 34229,
06304 Nice Cedex 4, France
9 Pontificia Universidad Catolica de Chile, Instituto de Astrofisica, Casilla 306, Santiago 22, Chile
Received 11 January 2016 / Accepted 31 March 2016
ABSTRACT
Context.The globular cluster HP 1 is projected at only 3.33 from the Galactic center. Together with its distance, this makes it one of the most central globular clusters in the Milky Way. It has a blue horizontal branch (BHB) and a metallicity of [Fe/H] ⇡ 1.0. This means that it probably is one of the oldest objects in the Galaxy. Abundance ratios can reveal the nucleosynthesis pattern of the first stars as well as the early chemical enrichment and early formation of stellar populations.
Aims.High-resolution spectra obtained for six stars were analyzed to derive the abundances of the light elements C, N, O, Na, and Al, the alpha-elements Mg, Si, Ca, and Ti, and the heavy elements Sr, Y, Zr, Ba, La, and Eu.
Methods.High-resolution spectra of six red giants that are confirmed members of the bulge globular cluster HP 1 were obtained with the 8 m VLT UT2-Kueyen telescope with the UVES spectrograph in FLAMES-UVES configuration. The spectroscopic parameter derivation was based on the excitation and ionization equilibrium of Fe
i
and Feii
.Results.We confirm a mean metallicity of [Fe/H] = 1.06 ± 0.10, by adding the two stars that were previously analyzed in HP 1. The alpha-elements O and Mg are enhanced by about +0.3 <⇠ [O,Mg/Fe] <⇠ +0.5 dex, Si is moderately enhanced with +0.15 <⇠ [Si/Fe] <⇠ +0.35 dex, while Ca and Ti show lower values of 0.04 <⇠ [Ca,Ti/Fe] <⇠ +0.28 dex. The r-element Eu is also enhanced with [Eu/Fe] ⇡ +0.4, which together with O and Mg is indicative of early enrichment by type II supernovae. Na and Al are low, but it is unclear if Na-O are anticorrelated. The heavy elements are moderately enhanced, with 0.20 < [La/Fe] < +0.43 dex and 0.0 < [Ba/Fe] < +0.75 dex, which is compatible with r-process formation. The spread in Y, Zr, Ba, and La abundances, on the other hand, appears to be compatible with the spinstar scenario or other additional mechanisms such as the weak r-process.
Key words. stars: abundances – Galaxy: bulge – globular clusters: individual: HP 1
1. Introduction
The globular cluster HP 1 is located at 3.33 and 1.8 kpc from the Galactic center. This is in the inner bulge volume, and the clus-ter lies among the globular clusclus-ters that are closest to the Galac-tic center. A metallicity of [Fe/H] ⇠ 1.0 was deduced from its color-magnitude diagram (CMD) by Ortolani et al. (1997,2011). High-resolution spectroscopy of two stars performed by Barbuy et al. (2006) resulted in [Fe/H] = 1.0 ± 0.2, and low-resolution spectroscopy of eight red giants by Dias et al. (2016) yielded a mean of [Fe/H] = 1.17 ± 0.07. Metallicities like this are at the lower end of the metal-poor stellar population in the metallic-ity distribution function (MDF) of bulge stars by Zoccali et al. (2008), Hill et al. (2011) and Rojas-Arriagada et al. (2014). A lower metallicity end at [Fe/H] ⇠ 1.0 is due to the fast chemi-cal enrichment in the Galactic bulge as modeled for example by
? Observations collected at the European Southern Observatory,
Paranal, Chile (ESO), under programs 93.D-0124A, 65.L-0340A, and 172.B-2002.
Cescutti et al. (2008). There are traces of a very metal-poor pop-ulation in the bulge, such as those found by García-Perez et al. (2013) and Howes et al. (2014,2015). These stars are very inter-esting, but most of them are located in the outer bulge and might more probably be halo stars. The main fact is that the bulk of the bulge stars show a lower end at [Fe/H] ⇠ 1.0. A metallic-ity of [Fe/H] ⇠ 1.0 could correspond to the population C or D as defined by Ness et al. (2013) in an MDF of a large sam-ple of bulge stars. For a latitude b = 5 , Ness et al. found a mean [Fe/H] = 0.66 for population C stars and identified it with the thick-disk population; their population D has a mean [Fe/H] = 1.16 and was identified by them as a metal-weak thick-disk population. It is not clear whether HP 1 fits into these categories.
The following evidence shows that there may be a stellar population peak at [Fe/H] ⇠ 1.0 in the bulge: (a) the metal-licity distribution of bulge globular clusters was shown to have two peaks at [Fe/H] ⇡ 0.5 and [Fe/H] ⇡ 1.0 (Bica et al.2016),
where a list of known bulge clusters was selected. This had al-ready been pointed out in Barbuy et al. (2006,2007,2009) and Rossi et al. (2015). (b) Another piece of evidence that stars of this metallicity are very old and are characteristic of the old Galactic bulge are the findings by Walker & Terndrup (1991), who showed that RR Lyrae in the bulge show a metallicity peak at [Fe/H] ⇡ 1.0 (see also Lee1992). Based on the MA-CHO survey, Kunder & Chaboyer (2008) determined a mean [Fe/H] = 1.25, with a broad metallicity distribution. Dékány et al. (2013) used the VISTA Variables in the Via Lactea (VVV) survey and obtained a spheroidal and centrally concentrated dis-tribution, with a slight elongation in its center. Using OGLE-III data, Pietrukowicz et al. (2012) derived [Fe/H] = 1.02 ± 0.18, with a barred distribution toward the central parts of the Galaxy. From OGLE-IV Pietrukowicz et al. (2015) obtained a mean [Fe/H] = 1.025 ± 0.25 and a triaxial ellipsoid shape. The outer bulge studied with 103 VVV RR Lyrae by Gran et al. (2016)
shows a centrally concentrated spheroidal distribution. (c) More recently, Schultheis et al. (2015) found a peak of bulge field stars at [Fe/H] ⇠ 1.0 that was enhanced in alpha-elements. (d) Schiavon et al. (in prep.) identified a sample of nitrogen-rich stars that also show metallicities of [Fe/H] ⇠ 1.0.
The triaxial ellipsoid shape as well as a cylindrical rotation as found in the bulge radial velocity assay (BRAVA) by Kunder et al. (2012) are expected from the dynamical evolution of an initially small classical bulge and a bar formed later in the disk (Saha et al.2012). Therefore either a spheroidal or triaxial shape would be consistent with an initially small classical bulge.
In summary, the moderately metal-poor globular clusters in the inner Galactic bulge might be relics of an early generation of long-lived stars formed in the proto-Galaxy. For this reason we have been pursuing the study of these clusters based on spec-troscopy (Barbuy et al.2006,2007,2009,2014), as well as pho-tometry and color-magnitude diagrams (CMDs) corrected for proper motion, as can be found in Ortolani et al. (2011) and Rossi et al. (2015).
In the present work we carry out a detailed analysis with high spectral resolution of six stars of the globular cluster HP 1. This cluster has a blue horizontal branch (BHB) combined with a metallicity of [Fe/H] ⇠ 1.0, which indicates that it is very old. Ortolani et al. (2011) derived the age by plotting HP 1 in the diagram of Fig. 17 by Dotter et al. (2010). From this, they computed an age di↵erence of about 1 Gyr for HP 1 compared to their sample of halo clusters with ⇠12.7 Gyr. This means that HP 1 is about 13.7 Gyr old and appears to be one of the oldest globular clusters in the Galaxy.
The cluster HP 1 is located at J2000 ↵ = 17h31m05.2s, =
29 5805400, with Galactic coordinates l = 2.58, b = +2.12.
The globular cluster HP 1 was discovered at the Observatoire de Haute Provence by Dufay et al. (1954). It was first studied through CMDs by Ortolani et al. (1997) by means of V, I col-ors. Davidge (2000) studied individual stars in the J, H, K, and CO filters and estimated [Fe/H] = 1.6 for HP 1. Ortolani et al. (2011) used J, H, and K with the multiconjugate adaptive op-tics demonstrator (MAD) at the VLT, applying a proper motion decontamination procedure, making use of the time di↵erence between the NTT observations from 1994 and the VLT/MAD observations in 2008. This allowed producing decontaminated CMDs and computing the orbit of HP 1 in the Galaxy. The CMD proper motion cleaning greatly optimizes the selection of member stars. Ortolani et al. (2011) showed that HP 1 re-mains confined within the bulge and/or bar. Minniti (1995) em-ployed medium-resolution infrared spectroscopy and measured indices in six stars of HP 1. They obtained a metallicity of
[Fe/H] = 0.56 and a radial velocity of 60 km s 1. Stephens
et al. (2004) used medium-resolution infrared spectra of six stars in HP 1 and derived a metallicity of [Fe/H] = 1.30. Two stars of HP 1 were observed at high spectral resolution with UVES and were analyzed spectroscopically (Barbuy et al.2006), pro-gram 65.L-0340 (PI: D. Minniti). To provide the most accu-rate information on the abundance pattern and kinematics of the metal-poor bulge globular clusters, we analyze a more signifi-cant number of stars in HP 1 to better identify the characteristics of the probably oldest stellar population in the Galaxy.
In this work we present a detailed abundance analysis using data from the FLAMES-UVES spectrograph at the VLT with a resolution R ⇠ 45 000 and a signal-to-noise ratio S/N > 200 for all sample stars. The MARCS model atmospheres are employed (Gustafsson et al.2008).
The observations are described in Sect. 2. The photometric e↵ective temperature and gravity are derived in Sect. 3. Spec-troscopic parameters are derived in Sect. 4 and abundance ratios are computed in Sect. 5. A discussion is presented in Sect. 6 and conclusions are drawn in Sect. 7.
2. Observations
The sample member stars of HP 1 were selected from data cor-rected for the proper motion that were reported in Ortolani et al. (2011). Figure1 shows a JHKs-combined image of HP 1 from
the Vista Variables in the Via Lactea VVV survey (Saito et al. 2012)1. The location of the six sample stars together with the two stars previously analyzed in Barbuy et al. (2006) are shown in Fig.2, with the observed field in a z-color image from the VVV survey.
The spectra of individual stars of HP 1 were obtained at the VLT using the UVES spectrograph (Dekker et al.2000) in FLAMES-UVES mode. The red arm (5800 6800 Å) has the ESO CCD # 20 chip, an MIT backside illuminated, with a size of 4096 ⇥ 2048 pixels, and a pixel size of 15 ⇥ 15 µm. The blue arm (4800 5800 Å) uses the ESO Marlene EEV CCD#44 chip, backside illuminated, with a size of 4102 ⇥ 2048 pixels, and a pixel size of 15 ⇥ 15 µm. The UVES standard setup 580 yields a resolution R ⇠ 45 000 for a slit width of 1 arcsec. The pixel scale is 0.0147 Å/pix, with ⇠7.5 pixels per resolution element at 6000 Å. The data were reduced using the UVES pipeline within the ESO/Reflex software (Ballester et al.2000; Modigliani et al. 2004). The log of the 2014 observations is given in Table 1. The spectra were flat fielded, optimally extracted, and wave-length calibrated with the FLAMES-UVES pipeline. Spectra ex-tracted from di↵erent frames were then co-added, taking into ac-count the radial velocities reported in Table2. The present UVES observations centered on 5800 Å yield a spectral coverage of 4800 < < 6800 Å, with a gap at 5708 5825 Å.
We measured the radial velocities of each run using the IRAF FXCOR cross-correlation method, as reported in Table2. The mean errors reported from the IRAF routine are 1.52 km s 1.
These values are clearly higher than jitter variations, which are estimated to be around 0.1 km s 1(e.g., Hekker et al.2008).
The present mean heliocentric radial velocity vhel
r = +40.0 ±
0.5 km s 1agrees very well with the value of +45.8 km s 1
de-rived from two stars in Barbuy et al. (2006). These are very low radial velocities, which combined with the low proper motions (Ortolani et al.2011) support the possibility of a confinement in the bulge.
Table 1. Log of the spectroscopic observations: dates, Julian dates, exposure times, airmass, seeing, and run number.
Date UT HJD Exp. Airmass Seeing Run
(s) (00) 07.06.14 04:33:15 2 456 815.69535 3000 1.007–1.013 0.72–0.84 1 24.06.14 01:55:50 2 456 832.58620 2700 1.083–1.026 0.47–0.73 2 20.07.14 05:13:40 2 456 858.72267 2700 1.329–1.591 0.52–0.53 3 30.07.14 01:26:09 2 456 868.56407 2700 1.004–1.023 0.83–0.75 4 14.08.14 01:00:33 2 456 883.54517 2750 1.010–1.047 0.71–0.69 5 14.08.14 01:58:18 2 456 883.54517 2750 1.048–1.126 0.64–0.72 6 14.08.14 02:45:54 2 456 883.61832 2750 1.129–1.264 0.77–1.00 7
Fig. 1.HP 1 JHKs-combined colour image from the VVV survey. The
image has size of 2 ⇥ 2 arcmin2. North is at 45 anticlockwise.
Figure3shows the V, I CMD of HP 1 and the location of the program stars on the red giant branch.
3. Photometric stellar parameters
3.1. Temperatures
The selected stars, their ID and 2MASS designations, coor-dinates, V, I magnitudes from Ortolani et al. (1997), and the 2MASS JHKs(Skrutskie et al.2006)2 and VVV JHKs
magni-tudes (Saito et al.2012) are listed in Table3. These magnitudes and colors were used to derive initial estimates of the e↵ective temperature and gravity, which were fine-tuned with spectro-scopic data using the Fe
i
and Feii
lines (see Sect. 4).For star 2461, the identification in the VVV images seems to indicate a blend of at least three stars, as can be seen in Fig.4. The VVV reductions were made in aperture photometry, which does not yield reliable magnitudes in this case. Only PSF pho-tometry, which is being carried out by the VVV group, will al-low distinguishing such cases. For the spectroscopy, the size of 1 arcsec of the fiber does allow observing the correct star.
The reddening value of E(B V) = 1.19 was derived from V, I CMDs by Ortolani et al. (1997). Barbuy et al. (1998) adopted
2 http://ipac.caltech.edu/2mass/releases/allsky/ 12.0 10.0 08.0 06.0 04.0 02.0 17:31:00.0 57:30.0 -29:58:00.0 30.0 59:00.0 30.0 -30:00:00.0 Right ascension Declination HP1 hp1-1 hp1-3 2939 2461 5485 5037 3514 2115
Fig. 2.z image from the VVV survey, indicating the location of the
sample stars.
Fig. 3.V, I CMD of HP 1 based on data from Ortolani et al. (1997), with the V, V I of the sample stars indicated.
Table 2. Radial velocities of the UVES sample stars in each of the seven exposure runs, corresponding heliocentric radial velocities, and mean heliocentric radial velocity.
Target vobsr vhel.r Target vobsr vhel.r km s 1 km s1 km s 1 km s1 2115 1 +37.024 +41.004 2461 1 +36.430 +40.41 2115 2 +45.313 +41.143 2461 2 +44.920 +40.75 2115 3 +57.650 +41.140 2461 3 +57.695 +41.185 2115 4 +60.847 +40.797 2461 4 +61.178 +41.128 2115 5 +65.612 +40.812 2461 5 +65.902 +41.102 2115 6 +65.699 +40.799 2461 6 +65.995 +41.095 2115 7 +65.819 +40.839 2461 7 +66.137 +41.157 Mean – +41.006 Mean – +41.048 2939 1 +41.756 +45.736 3514 1 +31.695 +35.675 2939 2 +50.289 +46.119 3514 2 +39.904 +35.734 2939 3 +62.512 +46.002 3514 3 +52.524 +36.014 2939 4 +66.057 +46.007 3514 4 +55.573 +35.523 2939 5 +70.718 +45.918 3514 5 +59.992 +35.192 2939 6 +70.859 +45.959 3514 6 +60.104 +35.204 2939 7 +70.974 +45.994 3514 7 +60.168 +35.188 Mean +46.035 Mean – +35.577 5037 1 +37.355 +41.335 5485 1 +30.851 +34.831 5037 2 +45.627 +41.457 5485 2 +38.234 +34.064 5037 3 +58.229 +41.719 5485 3 +50.745 +34.235 5037 4 +61.994 +41.994 5485 4 +54.665 +34.615 5037 5 +66.700 +41.900 5485 5 +59.565 +34.765 5037 6 +66.684 +41.784 5485 6 +59.782 +34.882 5037 7 +66.978 +41.998 5485 7 +59.832 +34.852 Mean +41.807 Mean – +34.679
E(B V) = 1.21 by using a slightly di↵erent RVvalue. Barbuy
et al. (2006) reported previous literature extinction values, and based on a reanalysis of CMD data, adopted E(B V) = 1.12. Ortolani et al. (2011) adopted E(V K) = 3.33, which translates into E(B V) = 1.21 using E(V K)/E(B V) = 2.744 (Rieke & Lebofsky1985); the use of a di↵erent reddening law does not a↵ect the result much. In the present work we adopted E(B V) = 1.12, following Barbuy et al. (2006).
E↵ective temperatures were derived from V I, V K, and J K using the color-temperature calibrations of Alonso et al. (1999, hereafter AAM99). These relations are similar to those by Ramirez & Meléndez (2005), as shown in their Fig. 11. The advantage of using the calibrations of Alonso et al. is that the bolometric corrections can be calculated, despite the disadvan-tage of having to translate Cousins to Johnson I, and 2MASS to TCS JHK colors. To translate V I from the Cousins to the Johnson system, we adopted (V I)C =0.778(V I)J(Bessell
1979). The J, H, KS2MASS magnitudes and colors were
trans-lated from the 2MASS system to CIT (California Institute of Technology), and from this to TCS (Telescopio Carlos Sánchez), using the relations from Alonso et al. (1998). The VVV JHKs
colors were translated into the 2MASS JHKssystem, using
re-lations reported by Soto et al. (2013). The derived photometric e↵ective temperatures are listed in Table4.
3.2. Gravities
The gravity values were computed using the classical formula log g⇤=4.44 + 4 logTT⇤ +0.4(Mbol⇤ Mbol ) + log MM ·⇤
We adopted T = 5770 K and Mbol = 4.75 for the Sun and
M⇤=0.85 M for the red giant branch (RGB) stars. For HP 1 we
Fig. 4.Image of star 2461 in the VVV survey, indicating a blend of
stars. Extraction of ⇠20 arcsec. North is 45 anticlockwise.
assumed a distance modulus of (m M)0=14.15 (Ortolani et al.
1997,2011), and the gravities were computed and are reported in Table4. Bolometric corrections were computed with formulae by AAM99. The photometric gravity values were computed as-suming the bolometric corrections from the temperature T(V I) values (Col. 9), and also with the spectroscopic temperature val-ues (Col. 11). The final spectroscopic gravities, described in the next section, are given in Col. 12 of Table4, and are compatible with the photometric gravities.
4. Spectroscopic stellar parameters
The equivalent widths (EW) were measured using the automatic code DAOSPEC, developed by Stetson & Pancino (2008). We also measured EWs line by line using IRAF for a few lines, in particular for Fe
ii
. The EWs measured for the Fei
and Feii
lines are reported in TableA.1. We limited EWs to 20 < EW(mÅ) < 99 to avoid on one hand, too weak lines that are a↵ected by the continuum level choice, in particular when using software such as DAOSPEC, and on the other hand, the saturated lines that are less sensitive to abundance variations.In Fig. 5we show the computed Fe
ii
lines compared with the observed spectra. Given the blends in lines Feii
6084.11 and 6456.39 Å, these lines were not used.In the line list given in Table A.1, literature oscillator strengths for Fe
i
from NIST3 and VALD34 databases (Martin et al.2002; Piskunov et al.1995) are reported. We also give the adopted values, where we have given preference to NIST over VALD. For Feii
we report the log g f values from Fuhr & Wiese (2006) and Meléndez & Barbuy (2009) in TableA.1. We adopted the latter values.Photospheric 1D models for the sample giants were ex-tracted from the MARCS model atmosphere grid (Gustafsson et al. 2008). We adopted the spherical and mildly CN-cycled set ([C/Fe] = 0.13, [N/Fe] = +0.31). This choice is due to the well-known mixing that occurs along the RGB, which transforms C into N. These models consider [↵/Fe] = +0.20 for [Fe/H] = 0.50 and [↵/Fe] = +0.40 for [Fe/H] 1.00. The LTE abundance analysis and the spectrum synthesis cal-culations were performed using the code described in Barbuy et al. (2003) and Coelho et al. (2005). An Fe abundance of
3 http://physics.nist.gov/PhysRefData/ASD/lines_form.
html
Table 3. Identifications, coordinates, V, I magnitudes from Ortolani et al. (1997), and JHKsmagnitudes from the 2MASS and VVV surveys. ID no. 2MASS ID ↵2000 2000 V I J H Ks JVVV HVVV KVVV 2115 – 17:31:05.420 29:59:20.71 17.070 14.281 – – – – – – 2461 17310822–2959108 17:31:08.22 29:59:10.85 17.115 14.628 – – – 12.066a 10.249a 10.959a 2939 17310729–2959021 17:31:07.31 29:59:02.30 17.131 14.394 11.901 10.869 10.595 – – – 3514 17310273–2958544 17:31:02.74 29:58:54.74 17.327 14.721 12.447 11.574 11.212 – 11.248 11.378 5037 17310347–2958259 17:31:03.48 29:58:26.04 17.101 14.487 12.302 11.303 10.994 12.287 11.502 11.115 5485 17310817–2958147 17:31:08.19 29:58:14.81 17.116 14.710 12.761 11.965 11.364 12.737 11.989 11.729 hp1-2 17310585–2958354 17:31:05.852 29:58:35.5 16.982 14.332 12.21 11.268 10.969 12.1527 11.4164 11.0541 hp1-3 17310587–2959180 17:31:05.872 29:59:18.1 17.003 14.358 12.167 11.146 10.828 12.1220 10.9702 10.9119 Notes. The two stars analyzed in Barbuy et al. (2006) are also listed.(a)Identification of star 2461 in VVV may be a blend of stars.
Table 4. Photometric stellar parameters derived using the calibrations by Alonso et al. (1999) for V I, V K, J K, bolometric corrections computed by adopting the V, I derived temperature, bolometric magnitudes, and corresponding gravity log g, and final spectroscopic parameters.
Photometric parameters Spectroscopic parameters
star TV I TV K TJ K TV K TJ K BCV Mbol log g Te↵ BCV log g log g [FeI/H] [FeII/H] [Fe/H] vt
2MASS VVV (K) (K) (K) (K) (K) (K) km s1 2115 4196.6 – – – – 0.702 0.18 1.95 4530 0.457 1.99 2.00 0.98 1.02 1.00 1.45 2461 4735.2 – – 4217.9 4506.5 0.368 0.14 2.04 4780 0.323 2.04 2.05 1.13 1.09 1.11 1.90 2939 4274.4 4025.1 4423.3 – – 0.638 0.12 1.98 4525 0.440 2.01 2.00 1.07 1.07 1.07 1.55 3514 4496.6 4249.4 4624.0 4855.9 – 0.487 0.08 2.09 4590 0.349 2.06 2.00 1.18 1.19 1.18 1.90 5037 4481.9 4254.3 4418.0 4324.0 4576.4 0.496 0.15 1.99 4570 0.429 2.05 2.15 0.98 1.03 1.0 1.20 5485 4920.4 4499.6 4197.0 4812.7 5152.4 0.298 0.14 2.08 4920 0.282 2.09 2.07 1.18 1.18 1.18 1.80 Notes. A reddening of E(B V) = 1.12 is assumed. For each star the last columns give the spectroscopic parameters from UVES spectra (present work).
FeII 5991.376 FeII 6084.110
FeII 6149.26 FeII 6247.56
FeII 6456.39 FeII 6516.08
Fig. 5.FeII lines for star 5037. Dotted lines: observed spectra. Blue
solid lines: synthetic spectra.
✏(Fe) = 7.50 (Grevesse & Sauval1998) was adopted. Molecular lines of MgH (A2⇧-X2⌃), CN (A2⇧-X2⌃), C
2Swan (A3⇧-X3⇧),
TiO (A3 -X3 ) and TiO (B3⇧-X3 ) ’ systems are taken into
account.
The stellar parameters were derived by initially adopting the photometric e↵ective temperature and gravity, and then by further constraining the temperature by imposing excitation
Fig. 6.Excitation and ionization equilibria of Fe
i
and Feii
lines for star 2939.equilibrium for Fe
i
lines. Fei
and Feii
lines allowed deriving gravities by imposing ionization equilibrium. Microturbulence velocities vtwere determined by canceling the trend of Fei
abun-dance vs. equivalent width.
The final spectroscopic parameters Te↵, log g, [Fe
i
/H],[Fe
ii
/H], [Fe/H], and vt values are reported in the last columnsof Table4. An example of excitation and ionization equilibria using Fe
i
and Feii
lines is shown in Fig.6for star HP1-2939.Table 5. Carbon, nitrogen, and oxygen abundances derived from C2
(0,1), CN (5,1), and [OI] lines.
Line (Å) 2115 2641 2939 3514 5037 5485 C2(0,1) 5635.5 +0.0 +0.0 +0.0 +0.0 +0.0 +0.0
CN(5,1) 6332.2 +0.7 +0.5 +0.5 +0.8 +0.5 +0.5 [OI] 6300.3 +0.4 +0.5 +0.5 +0.4 +0.35 +0.4
5. Abundance ratios
Abundance ratios were obtained by means of line-by-line spec-trum synthesis calculations compared to the observed spectra for the line lists given in Tables5,B.1–B.3. The fits were made by eye using a series of calculations for di↵erent abundances and verifying the continua over 20 Å. We then zoomed and fine-tuned them locally. The solar abundances were adopted from Grevesse et al. (1998), which are close to the latest values by As-plund et al. (2009), Grevesse et al. (2014) or Lodders (2009). For oxygen we adopted ✏(O) = 8.77 following Allende Prieto et al. (2001) for 1D model atmospheres, which is very close to the value of ✏(O) = 8.76 recommended from 3D models by Ste↵en et al. (2015).
5.1. Carbon, nitrogen, and oxygen
The carbon abundances were estimated from the C2(0,1)
band-head at 5635.3 Å. The list of laboratory C2 lines (Phillips &
Davis1968) was reported in Barbuy et al. (2014). Because the feature is weak in these metal-poor stars, we adopted [C/Fe] = 0.0 and just checked that this was a suitable upper limit. The ni-trogen abundances were measured using the CN (5, 1) 6332.18 Å of the CN A2⇧-X2⌃red system. The CN feature is also weak,
but can be used to estimate a reliable upper limit for the N abundance.
The forbidden oxygen [OI] 6300.311 Å line was used to de-rive the oxygen abundances. A check for telluric lines was car-ried out by overplotting the spectrum of a fast-rotation B star, and we verified that none of the stars were a↵ected by them. The resulting C, N, and O abundances are given in Table5, show-ing the essentially adopted carbon abundance of [C/Fe] ⇠ 0.0, the enhanced nitrogen abundances as expected in giants, and en-hanced oxygen abundances of +0.35 < [O/Fe] < +0.5, typical of enrichment by SNII.
5.2. Odd-Z elements Na, Al, and alpha-elements
In TablesB.1andB.2we report the line-by-line abundances of the odd-Z elements Na, Al, and the alpha-elements Mg, Si, Ca, and Ti abundances. We have inspected the abundance results as a function of e↵ective temperature and microturbulence velocity and found no trend for any of the elements. In addition to the six sample stars, we also rederived the abundances of Na, Al and Mg for the two stars analyzed in Barbuy et al. (2006), as given in Table 6.
5.3. Heavy elements
In TableB.3we report the line-by-line derivation of abundances for lines of the neutron-capture dominant s-elements Sr, Y, Zr, La, Ba, and the r-element Eu. As in TableB.1, we also rederived the abundances of heavy elements for the two stars analyzed in Barbuy et al. (2006). The hyperfine structure (HFS) for the studied lines of La
ii
, Baii
and Euii
were taken into account,Table 6. Abundance uncertainties for star 2115 for uncertainties of Te↵ = +100 K, log g = +0.2, vt = 0.2 km s 1, an assumed er-ror in EWs or continuum placement, and the corresponding total erer-ror.
Abundance T log g vt EW (Px2)1/2 100 K 0.2 dex 0.2 kms1 (1) (2) (3) (4) (5) [FeI/H] +0.06 +0.03 0.07 0.02 0.10 [FeII/H] 0.12 +0.11 0.03 0.02 0.17 [N/Fe] 0.05 +0.00 +0.00 0.10 0.11 [O/Fe] +0.00 0.05 +0.00 0.05 0.07 [NaI/Fe] +0.10 +0.00 0.01 0.10 0.14 [Al/Fe] +0.05 +0.00 +0.00 0.10 0.11 [MgI/Fe] +0.03 +0.00 +0.00 0.10 0.10 [SiI/Fe] +0.00 +0.02 0.02 0.05 0.06 [CaI/Fe] +0.10 +0.00 0.05 0.05 0.12 [TiI/Fe] +0.10 +0.00 0.03 0.05 0.12 [TiII/Fe] 0.10 +0.08 0.01 0.05 0.14 [SrI/Fe] 0.15 +0.00 +0.00 0.20 0.25 [YI/Fe] 0.15 +0.00 +0.00 0.10 0.18 [YII/Fe] +0.02 0.15 +0.00 0.10 0.18 [ZrI/Fe] 0.20 +0.00 +0.00 0.10 0.22 [BaII/Fe] +0.02 +0.02 0.20 0.10 0.23 [LaII/Fe] +0.00 +0.03 +0.00 0.10 0.10 [EuII/Fe] +0.00 +0.03 +0.00 0.10 0.10 Notes. The errors correspond to the di↵erence in abundance obtained with the modified parameters.
as described in Barbuy et al. (2014). The fits to the lines of Eu
ii
6645, Baii
6141 and Baii
6496 Å lines in star 2115 are shown in Fig.8.5.4. Errors
The errors due to uncertainties in spectroscopic parameters are given in Table6, applied to the sample star HP 1: 2115. The error on the slope in the FeI vs. excitation potential implies an error in the temperature of ±100 K for the sample stars. An uncertainty of the order of 0.2 km s 1 on the microturbulence velocity is
estimated from the imposition of a constant value of [Fe/H] as a function of EWs. Errors based on EWs are given on FeI and FeII abundances.
The errors on the abundance ratios [X/Fe] were computed by fitting the lines with the modified atmospheric model. The error reported corresponds to the new value obtained by us-ing the modified model atmosphere. The element abundance ra-tios, induced by a change of Te↵ = +100 K, log g = +0.2,
vt=0.2 km s 1, and a total error estimate, are given in Table6. These errors are overestimated because the stellar parameters are covariant. The correlation matrix is difficult to estimate, how-ever, and would add other error sources, so that we preferred the quadratic sum of the diagonal terms as reliable.
Additionally, an uncertainty of about 0.8 mÅ in the EWs of the Fe lines was estimated with the formula of (Cayrel 1988, 2004). With a mean FWHM = 12.5 pixels, or 0.184 Å, a CCD pixel size of 15 µm, or x = 0.0147 Å in the spectra, and assuming a mean S/N = 100, we obtain an error EW ⇠ 0.8 mÅ.
We derived abundances uniquely from fitting synthetic spec-tra, such that we need to take the uncertainty in the continuum placement into account. We estimate an error of 0.1 dex for the weak and strong lines and 0.05 dex for medium lines. This is included in Table6.
Table 7. Mean abundances of C, N, odd-Z elements Na, Al, ↵-elements O, Mg, Si, Ca, Ti, and heavy elements Y, Sr, Zr, Ba, La, and Eu. [X/Fe] 2115 2461 2939 3514 5037 5485 HP 1-2 HP 1-3 Mean NGC 6522 C +0.00 +0.00 +0.00 +0.00 +0.00 +0.00 +0.00 +0.00 +0.00 0.03 N +0.70 +0.50 +0.50 +0.80 +0.50 +0.50 +0.50 +0.20 +0.53 +0.67 O +0.40 +0.50 +0.50 +0.40 +0.35 +0.40 +0.30 +0.30 +0.40 +0.36 Na +0.03 0.30 0.23 0.10 0.10 0.30 0.05 0.20 0.16 +0.05 Al 0.15 +0.05 +0.05 +0.08 +0.20 +0.20 0.28 +0.18 +0.04 +0.20 Mg +0.30 +0.35 +0.65 +0.35 +0.33 +0.40 +0.15 +0.33 +0.36 +0.23 Si +0.25 +0.31 +0.33 +0.31 +0.20 +0.15 +0.30 +0.30 +0.27 +0.13 Ca +0.28 +0.11 +0.21 +0.11 +0.26 +0.04 0.04 +0.10 +013 +0.13 TiI +0.26 +0.16 +0.28 +0.17 +0.16 +0.06 +0.07 +0.08 +0.16 +0.04 TiII +0.40 +0.18 +0.23 +0.19 +0.31 +0.13 +0.10 +0.15 +0.21 +0.17 Y +0.15 +0.30 +0.25 +0.35 +0.10 +0.50 -0.15 0.1 +0.18 +0.31 Zr +0.15 0.10 +0.28 +0.37 +0.20 +0.15 – 0.02 +0.15 +0.18 Sr +0.55: – +0.30 – +0.40 – +0.50: +0.30: +0.41: +0.23 Ba +0.50 +0.30 +0.50 +0.00 +0.75 +0.45 +0.65 +0.47 +0.46 +0.32 La +0.33 +0.23 +0.43 +0.32 +0.08 +0.10 0.15 0.18 +0.23 – Eu +0.50 +0.50 +0.70 +0.60 +0.60 +0.55 +0.40 +0.55 +0.15 +0.30 Ba/Eu 0.0 0.20 0.20 0.60 +0.15 0.10 +0.25 0.08 +0.31 – vr(km s 1) +41.0 +41.0 +46.0 +35.0 +41.0 +34.0 +44.6 +44.0
Notes. Mean values for NGC 6522 are given in the last column.
6. Discussion
In Table7we report the mean abundances for each star and for each element, and in the last line we report the radial velocities. 6.1. Two stellar populations?
We derived a mean metallicity of [Fe/H] = 1.06 ± 0.15 by adding the two stars previously analyzed in HP 1 by Barbuy et al. (2006).
Two stars, 3514 and 5485, show a slightly lower metallic-ity of [Fe/H] = 1.18 and, possibly not by coincidence, they also show radial velocities of vhel
r = +35.6 and +34.7 km s 1,
respectively, which is lower than the radial velocities of the other six stars, whose values lie in the range 41.0 < vhel
r <
47.0 km s 1. Their radial velocities are compatible with being
members within the uncertainties. If these two stars correspond to an earlier stellar generation in the cluster, then we might have a first stellar generation with [Fe/H] = 1.18 and a second gen-eration with a mean metallicity of [Fe/H] = 1.02 ± 0.05, as found for the other six stars.
The membership of these two lower velocity stars is another question. Recent work by Bellini et al. (2015) for NGC 2808 has shown di↵erences in radial velocities between two stellar popu-lations in a cluster. Malavolta et al. (2015) studied M4, showing that it has a dispersion of 4 km s 1. HP 1 is probably subject
to tidal e↵ects that lead to possible perturbations or even disrup-tion. The two low-velocity stars might become unbound from the cluster, given their di↵erence in velocites of vhel
r = 10 km s 1.
6.2. Odd-Z elements and Na-O anticorrelation
Na and Al tend to be underabundant in HP 1. This may be due to internal nucleosynthesis (Gratton et al.2012). In particular, a Na-O anticorrelation due to the Ne-Na cycle of proton capture reactions was shown by Carretta et al. (2009) to be confirmed for several globular clusters. In Fig.7we show the [Na/Fe] vs. [O/Fe] for 14 stars of NGC 6121 by Carretta et al. (2009), com-pared to the present results for HP 1. This figure points to evi-dence for a Na-O anticorrelation in HP 1, but with the di↵erence
Fig. 7.[Na/Fe] vs. [O/Fe] for the sample stars compared with stars of
NGC 6121. Symbols: present work: green filled squares: the 6 HP 1 stars with [Fe/H] = 1.0 ± 0.05; open green squares: the 2 HP 1 stars with [Fe/H] = 1.18; compared with blue filled triangles: stars of NGC 6121.
with respect to NGC 6121 and most other clusters studied by Carretta et al. (2009), of Na being lower than in the other clus-ters. Na is originally probably underabundant in HP 1, which reinforces its peculiar pattern.
6.3. Alpha-elements
The ↵-element enhancements in O, Mg, and Si together with the enhancement of the r-process element Eu are indicative of a fast early enrichment by SNII. Ca and Ti are only slightly enhanced with [Ca/Fe] = +0.13 and [Ti/Fe] = +0.18 (a mean of Ti
i
andFig. 8. Fits to the Eu
ii
6645, Baii
6141 and Baii
6496 Å lines in star 2115. The [Eu/Fe] and [Ba/Fe] values adopted for the di↵erent cal-culations are indicated in the panels.Ti
ii
abundances). The di↵erence between O, Mg, and Si on one hand and Ca and Ti on the other was also detected in other bulge samples, such as NGC 6522 (Barbuy et al. 2014). The mean values for NGC 6522 are reported for comparison purposes in Table 10.The fact that [O, Mg, Si/Fe] are higher than [Ca, Ti/Fe] does not match the results for field stars by González et al. (2011). This might be explained by the mass of supernovae type II that enriched the cluster, or else by the absence of contribution from supernovae type I.
6.4. Heavy elements
Figure9shows the results for Y, Zr, Ba, La, and Eu compared to other available heavy element abundance determinations in bulge stars. The Sr abundances are not very reliable because the lines used are extremely faint. We therefore did not plot this ele-ment. Literature data include a) field red giants in Plaut’s field analyzed by Johnson et al. (2012); b) seven red giants in the globular cluster M62 (NGC 6266) by Yong et al. (2014); c) mi-crolensed bulge dwarf stars analyzed by Bensby et al. (2013); d) four red giants in NGC 6522 analyzed by Barbuy et al. (2014); e) 56 bulge field red giants for which Van der Swaelmen et al. (2016) derived heavy element abundances; and f) five bulge field red giants with [Fe/H] ⇡ 1.0 from the sample by Ness et al. (2013) that were analyzed by Siqueira-Mello et al. (2016).
This figure shows that [Zr/Fe], [Ba/Fe], [La/Fe], and [Eu/Fe] increase steadily with decreasing metallicity. At a metallicity of around [Fe/H] ⇡ 1.0, the dominantly s-elements Y, Zr, Ba, and La show an abundance spread. This behavior is compatible with expectations from massive spinstars: a spread like this is pre-dicted from the models by Frischknecht et al. (2016, and refer-ences therein) and Meynet et al. (2016, and referrefer-ences therein), as discussed and shown in Chiappini et al. (2011), Chiappini (2013), and Barbuy et al. (2014).
On the other hand, the enhancements of O, Mg, and Eu, and to a lesser extent of Si, Ca, and Ti, indicate an early enrichment
by supernovae type II. The small fraction of r-process production of the dominantly s-elements in the solar neighborhood might be responsible for their production at early times, as first suggested by Truran (1981). To verify the r- or s-nature of the heavy el-ement abundances, we report the [Ba/Eu] ratios in penultimate line of Table7.
The ratio [Ba/Eu] is indicative of the relative contribution of the s- and r-process. We find 0.60 < [Ba/Eu] < +0.25, and [Ba/Eu]r ⇠ 0.8 is given as typical of a pure r-process by
Bisterzo et al. (2014). This means that the barium-to-europium ratio is above the line for a pure r-process and might be in-terpreted as due to an s-process. The s-process can be due to transfer of matter from an AGB companion or to a general AGB pollution, as recently made likely to occur in multi-population clusters (e.g. Renzini et al.2015), or to spinstars. The spread in abundances of Y, Sr, Zr, Ba, and La at around [Fe/H] ⇠ 1.0 is the main indicator of the contribution by spinstars (Chiappini et al.2011). Finally, both r- and s- processes may take place, that is, massive spinstars producing s-elements may not exclude a later explosion of the supernova and the subsequent r-process. Another possible explanation would be an additional process that enhanced the lightest heavy elements at early times. Several models are proposed in the literature, such as the lighter element primary process LEPP. Travaglio et al. (2004) discussed this pro-cess for the Sun. A di↵erent LEPP mechanism was discussed by Montes et al. (2007) for metal-poor stars. A possible explana-tion for the LEPP is the weak r-process (Wanajo & Ishimaru 2006) or supernovae neutrino-driving winds (Arcones & Thiele-mann2013). We note that di↵erent processes might be enriching in Y, Sr, and Zr at di↵erent metallicities and environments. Niu et al. (2015) proposed a unified solution. Finally, it is important to mention that, as pointed out by Roederer et al. (2010; see his Fig. 11), there are varying degrees of enrichment of first- and second-peak elements in stars enriched by neutron capture.
7. Conclusions
We carried out a detailed analysis of six red giants of the bulge moderately metal-poor globular cluster HP 1 and added two other previously analyzed stars. A metallicity of [Fe/H] = 1.06 ± 0.15 was derived from the eight stars, and since HP 1 has a blue horizontal branch, this combination of characteristics is an indication of a very old age.
We found overabundances of the ↵-elements [Mg/Fe] ⇡ [O/Fe] ⇡ +0.4 and [Si/Fe] ⇡ +0.3 and lower values of [Ca/Fe] ⇡ [Ti/Fe] ⇡ +0.10. The light odd-Z elements Na and Al are low with [Na/Fe] = 0.20 and [Al/Fe] = +0.18. In particular, be-cause of the low Na abundance, the Na-O anticorrelation is un-clear. HP 1 has a relatively low mass, with absolute magnitude MV = 6.46, therefore it is probably less prone to show this
e↵ect, as demonstrated for several clusters by Carretta et al. (2009).
The dominantly s-elements are moderately enhanced with [Ba/Fe] ⇡ +0.32 and the r-element [Eu/Fe] ⇡ +0.30. These val-ues are very similar to the results for NGC 6522 (Barbuy et al. 2014), indicating that these two clusters that are located in the central parts of the Galactic bulge may be of similar origin.
The abundances in the earliest globular clusters can reveal the nature of the first stars and also the location at which they formed in the Galaxy. It appears to be of great importance to further investigate this and other such clusters in terms of metal-licity, abundances, kinematics, orbits, and ages.
Fig. 9.[Y, Zr, Ba, La, Eu/Fe] vs. [Fe/H] and [Y/Ba] vs. [Fe/H] for the sample stars compared with literature values. Symbols: blue filled squares: NGC 6522; magenta pentagons: M62; red filled triangles: bulge field dwarfs by Bensby et al. (2013); blue triangles: bulge field red giants from Rich et al. (2012); cyan triangles: Van der Swaelmen et al. (2016); yellow triangles: Siqueira-Mello et al. (2016); green filled pentagons: HP 1 from the present work.
Acknowledgements. B.B., E.C., A.V., C.S.M., H.E., and E.B. acknowledge grants and fellowships from CNPq, Capes and Fapesp. S.O. acknowledges the financial support from the Università di Padova and from the Italian Ministero dell’Università e della Ricerca Scientifica e Tecnologica (MURST), Italy. D.M. and M.Z. acknowledge support from the BASAL Center for Astrophysics and Associated Technologies PFB-06, the Ministry of Economy, Development, and Tourism’s Millennium Science Initiative through grant IC120009, awarded to The Millennium Institute of Astrophysics (MAS), and Proyectos Fondecyt Reg-ular 1130196 and 1150345.
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Appendix A: Equivalent widths and atomic data of FeIand FeIIlines
Table A.1. Fe
i
and Feii
lines, their wavelengths, excitation potential (eV), oscillator strengths from NIST and VALD3 for Fei
and Fuhr & Wiese (2006) and Meléndez & Barbuy (2009) for Feii
, and equivalent widths (mÅ).Species (Å) ex(eV) log g fNIST log g fVALD log g fadopted 2115 2461 2939 3514 5037 5485
FW06 MB09 FeII 5991.38 3.15 3.65 3.54 3.54 – 33.8 – 24.6 27.5 26.2 FeII 6149.25 3.89 2.84 2.69 2.69 26.3 32.9 23.5 31.4 22.3 33.0 FeII 6247.56 3.89 2.43 2.30 2.30 24.1 38.9 27.8 – 28.0 46.1 FeII 6416.93 3.89 2.88 2.64 2.64 24.9 26.0 23.1 20.6 – 33.5 FeII 6432.68 2.89 3.50 3.57 3.57 33.6 41.7 32.4 34.1 29.2 45.5 FeII 6516.08 2.89 3.372 3.31 3.31 45.7 59.9 47.9 51.7 44.0 64.0 FeI 5525.544 4.231 – 1.084 1.084 – 37.7 – – – 40.5 FeI 5543.360 4.218 – 1.140 1.11 – 45.5 – – 62.1 35.2 FeI 5554.894 4.549 – 0.440 0.440 – 67.9 67.8 – 74.4 48.6 FeI 5560.211 4.435 – 1.190 1.16 39.4 29.5 41.3 – 34.0 22.5 FeI 5567.391 2.609 – 2.564 2.67 75.8 58.1 77.6 – 68.0 51.2 FeI 5618.632 4.209 – 1.276 1.276 38.3 30.5 49.5 – 34.1 34.6 FeI 5619.595 4.387 – 1.700 1.67 31.1 – 27.5 – 23.6 – FeI 5633.946 4.991 – 0.270 0.32 44.7 44.7 56.6 – 38.4 33.6 FeI 5638.262 4.220 – 0.870 0.84 63.4 62.2 76.1 – 60.6 34.9 FeI 5641.434 4.256 – 1.180 1.15 58.8 48.3 59.5 – 60.1 39.2 FeI 5653.865 4.387 – 1.640 1.61 26.2 – – – 20.9 – FeI 5679.023 4.652 – 0.920 0.90 – 31.8 34.5 – 42.5 36.4 FeI 5701.544 2.559 – 2.216 2.216 – – – – 94.0 70.7 FeI 5706.096 4.283 – 3.012 3.012 – 58.3 – – – – FeI 5717.833 4.284 – 1.130 1.10 28.8 28.7 36.2 46.4 42.9 37.0 FeI 6056.005 4.733 – 0.460 0.460 48.8 42.8 54.7 – 50.8 34.6 FeI 6078.491 4.796 – 0.321 0.321 59.3 49.5 52.0 55.9 56.5 44.3 FeI 6079.008 4.652 1.100 1.120 1.100 30.8 31.8 26.5 24.2 31.4 – FeI 6082.710 2.223 3.573 3.573 3.573 67.6 48.2 61.2 50.9 53.8 24.5 FeI 6093.643 4.608 1.470 1.599 1.470 20.8 – – – 24.2 – FeI 6096.664 3.984 1.880 1.930 1.88 34.4 23.8 28.9 23.1 30.5 – FeI 6151.617 2.176 3.299 3.299 3.299 80.9 76.2 81.1 78.2 76.2 47.5 FeI 6157.728 4.076 1.22 1.260 1.22 69.6 60.8 72.4 40.0 64.0 41.8 FeI 6165.360 4.143 1.474 1.474 1.474 53.1 44.2 47.3 49.0 49.7 42.2 FeI 6180.203 2.728 2.649 2.586 2.649 84.4 74.9 87.3 79.7 79.0 61.7 FeI 6187.989 3.943 1.67 1.720 1.67 51.7 39.9 56.7 46.5 43.2 23.3 FeI 6200.313 2.609 2.437 2.437 2.437 – – – 97.7 – – FeI 6219.281 2.198 2.433 2.433 2.43 – – – – – 97.9 FeI 6226.734 3.884 – 2.220 2.220 – – 21.6 – – – FeI 6229.226 2.845 2.805 2.805 2.805 26.8 36.9 51.7 45.6 41.5 22.8 FeI 6240.646 2.223 3.173 3.233 3.173 75.7 54.9 74.6 – 70.0 37.3 FeI 6246.318 3.603 0.877 0.733 0.877 – 68.6 – – 90.8 80.9 FeI 6265.132 2.176 2.550 2.550 2.550 – – – – – 95.5 FeI 6270.223 2.858 2.609 2.464 2.609 76.8 64.5 68.7 60.7 66.5 44.1 FeI 6271.278 3.332 2.703 2.703 2.703 24.8 20.8 24.9 – – – FeI 6302.494 3.686 – 0.973 0.973 84.7 77.9 83.7 – 85.0 62.2 FeI 6311.500 2.832 3.141 3.141 3.141 48.8 25.9 42.8 30.4 39.0 – FeI 6315.306 4.143 1.232 1.232 1.232 62.8 53.5 61.0 57.3 56.0 38.9 FeI 6315.811 4.076 1.66 1.710 1.66 40.0 – 40.0 33.0 36.8 24.4 FeI 6322.685 2.588 2.426 2.426 2.426 – 95.1 – 95.8 – 82.9 FeI 6336.823 3.686 0.856 0.856 0.856 – 98.3 – – – 88.8 FeI 6344.148 2.433 2.923 2.923 2.92 – 87.8 – 88.8 91.0 62.1 FeI 6355.028 2.845 2.291 2.350 2.291 88.6 71.3 79.8 – 78.3 51.7 FeI 6380.743 4.186 1.376 1.376 1.376 50.0 45.3 47.4 33.7 51.1 31.3 FeI 6392.538 2.279 – 4.030 4.030 42.5 21.6 38.0 24.3 32.6 – FeI 6408.017 3.686 1.018 1.018 1.018 96.7 86.0 87.8 84.4 85.4 77.1 FeI 6411.648 3.654 0.718 0.595 0.718 – – – – – 98.9 FeI 6419.949 4.733 0.27 0.240 0.27 66.7 58.8 65.4 63.8 66.4 47.0 FeI 6469.192 4.835 0.81 0.770 0.81 – 26.9 – – 38.5 – FeI 6475.624 2.559 2.942 2.942 2.942 76.4 64.5 81.4 68.6 83.4 50.4
Table A.1. continued.
Species (Å) ex(eV) log g fNIST log g fVALD log g fadopted 2115 2461 2939 3514 5037 5485
FeI 6481.870 2.279 2.984 2.984 2.984 96.4 79.4 98.3 89.2 – 69.9 FeI 6498.938 0.958 4.687 4.699 4.687 – 80.0 – 94.8 – 58.9 FeI 6518.366 2.832 2.298 2.460 2.298 82.0 68.5 86.8 – 76.6 53.1 FeI 6569.214 4.733 0.450 0.420 0.450 63.4 60.5 65.2 67.2 63.1 46.2 FeI 6574.226 0.990 5.004 5.023 5.004 93.5 62.3 89.5 78.8 80.8 44.5 FeI 6575.016 2.588 2.710 2.710 2.710 95.2 79.8 93.2 85.7 83.3 65.3 FeI 6581.209 1.485 4.679 4.679 4.679 56.8 28.3 58.4 39.7 42.2 21.7 FeI 6593.870 2.433 2.422 2.422 2.422 – – – – – 92.8 FeI 6597.559 4.796 1.050 1.070 1.05 29.5 20.6 26.4 31.7 26.3 – FeI 6608.025 2.279 – 4.030 4.030 39.3 – 38.0 27.9 35.9 – FeI 6609.110 2.559 2.692 2.692 2.692 95.4 77.9 93.0 87.7 82.7 61.8 FeI 6627.544 4.549 – 1.680 1.680 20.5 – – – – –
Appendix B: Resulting abundances
In Tables B.1–B.3 are listed line-by-line, the resulting abun-dance ratios [X/Fe] for the light elements Na, Mg, and Al,
Table B.1. Abundances of light elements Na, Mg, Al.
Species (Å) ex(eV) log g f 2115 2461 2939 3514 5037 5485 HP1-2 HP1-3
NaI 5682.633 2.102439 0.706 0.15 0.30 0.30 +0.00 0.10 0.30 +0.00 0.20 NaI 5688.194 2.104571 1.400 +0.20 0.30 0.15 -0.20 0.10 0.30 0.10 0.20 NaI 5688.205 2.104571 0.45 +0.20 0.30 0.15 -0.20 0.10 0.30 0.10 0.20 NaI 6154.230 2.102439 1.56 0.05 0.10 0.15 0.15: +0.20 0.30 0.30 +0.05 NaI 6160.753 2.104571 1.26 +0.25 0.15 0.15 +0.10 +0.30 +0.00 0.30 – AlI 6696.185 4.021753 1.576 0.1 +0.0 +0.1 +0.00 +0.20 +0.20: 0.25 +0.15 AlI 6696.788 4.021919 1.421 – – – – – – – – AlI 6698.673 3.142933 1.647 0.10 +0.10 0.00 +0.10 +0.20 +0.20 0.30 +0.20 MgI 6318.720 5.108171 2.10 +0.30 +0.30 +0.50 +0.40 +0.40 +0.40 +0.15 +0.35 MgI 6319.242 5.108171 2.36 +0.30 +0.40 +0.30 +0.30 +0.20 +0.40 – +0.30
the alpha-elements O, Mg, Si, Ca, and Ti, and the heavy elements Eu, Ba, La, Y, Zr, and Sr.
Table B.2. Abundance of alpha-elements O, Mg, Si, Ca, Ti.
Species (Å) ex(eV) log g f 2115 2461 2939 3514 5037 5485
SiI 5665.555 4.920417 2.04 +0.1 +0.2 +0.2 +0.3 0.0 +0.2 SiI 5666.690 5.616073 1.74 – – – +0.3 +0.1 +0.2 SiI 5690.425 4.929980 1.87 +0.3 +0.2 +0.3 +0.3 +0.2 +0.0 SiI 5948.545 5.082689 1.30 +0.4 +0.4 +0.5 +0.2 +0.3 +0.2 SiI 6142.494 5.619572 1.50 +0.3 +0.5 +0.25 – +0.2 +0.1 SiI 6145.020 5.616073 1.45 +0.3 +0.3 +0.2 – +0.1 +0.3 SiI 6155.142 5.619572 0.85 +0.2 +0.2 +0.4 +0.4 +0.2 – SiI 6237.328 5.613910 1.01 +0.2 +0.2 +0.3 +0.2 +0.2 +0.1 SiI 6243.823 5.616073 1.30 +0.2 – +0.3 +0.4 +0.1 – SiI 6414.987 5.871240 1.13 +0.3 +0.4 +0.4 +0.3 +0.4 +0.2 SiI 6721.844 5.862872 1.17 +0.2 +0.20 +0.4 +0.4 +0.2 +0.0 CaI 5601.277 2.525852 0.52 +0.0 +0.0 +0.0 0.1 +0.0 0.3 CaI 5867.562 2.932710 1.55 +0.1 +0.1 +0.0 +0.3 +0.0 – CaI 6102.723 1.879467 0.79 +0.3 +0.1 +0.1 0.2 +0.2 +0.0 CaI 6122.217 1.885935 0.20 +0.0 +0.1 +0.0 +0.0 +0.2 +0.1 CaI 6161.295 2.523157 1.02 +0.3 +0.1 +0.4 +0.2 +0.2 +0.0 CaI 6162.167 1.899063 0.09 +0.3 +0.1 +0.1 +0.1 +0.2 +0.0 CaI 6166.440 2.521433 0.90 +0.3 +0.0 +0.4 +0.1 +0.3 +0.0 CaI 6169.060 2.523157 0.54 +0.3 +0.1 +0.1 +0.2 +0.4 +0.1 CaI 6169.564 2.525852 0.27 +0.3 +0.1 +0.2 +0.0 +0.2 +0.0 CaI 6439.080 2.525852 +0.30 +0.3 +0.0 +0.1 +0.0 +0.4 +0.1 CaI 6455.605 2.523157 1.35 +0.3 +0.1 +0.4 +0.2 +0.4 +0.0 CaI 6464.679 2.525852 2.10 +0.4 +0.4 >+0.5 – +0.2 – CaI 6471.668 2.525852 0.59 +0.3 +0.2 +0.3 +0.2 +0.4 +0.3 CaI 6493.788 2.521433 0.00 +0.3 0.1 +0.1 +0.0 +0.3 +0.0 CaI 6499.654 2.523157 0.85 +0.3 +0.0 +0.3 +0.1 +0.2 +0.0 CaI 6572.779 0.00 4.32 +0.5 0.0 +0.3 +0.2 +0.3 +0.0 CaI 6717.687 2.709192 0.61 +0.4 +0.2 +0.5 +0.3 +0.5 +0.1 TiI 5689.459 2.296971 0.47 +0.2 – – +0.1 – – TiI 5866.449 1.066626 0.84 +0.1 +0.0 +0.2 +0.0 +0.3 – TiI 5922.108 1.046078 1.46 +0.3 +0.2 +0.4 +0.1 +0.0 – TiI 5941.750 1.052997 1.53 +0.2 +0.2 +0.3 +0.2 +0.2 +0.2 TiI 5965.825 1.879329 0.42 +0.3 +0.2 +0.3 +0.1 +0.1 +0.0 TiI 5978.539 1.873295 0.53 +0.3 +0.1 +0.2 +0.2 +0.2 +0.0 TiI 6064.623 1.046078 1.94 +0.3 +0.2 +0.4 +0.2 +0.3 – TiI 6091.169 2.267521 0.42 +0.3 +0.2 +0.3 +0.3 +0.0 +0.3 TiI 6126.214 1.066626 1.43 +0.3 +0.2 +0.3 +0.2 +0.2 +0.0 TiI 6258.110 1.443249 0.36 +0.1 +0.0 +0.3 +0.0 +0.1 0.1 TiI 6261.106 1.429852 0.48 +0.3 +0.1 +0.2 +0.0 +0.0 0.1 TiI 6303.767 1.443249 1.57 +0.3 +0.15 +0.35 +0.3 +0.2 +0.2 TiI 6312.240 1.460236 1.60 +0.3 +0.3 +0.4 +0.2 +0.1 – TiI 6336.113 1.443249 1.74 +0.1 – +0.2 +0.2 +0.0 – TiI 6508.150 1.429852 2.05 +0.2 +0.20 +0.2 +0.3 +0.2 – TiI 6554.238 1.443249 1.22 +0.1 +0.0 +0.2 +0.2 +0.1 0.15 TiI 6556.077 1.460236 1.07 +0.2 +0.2 +0.3 +0.1 +0.1 +0.0 TiI 6599.113 0.899612 2.09 +0.3 +0.2 +0.4 +0.2 +0.5 +0.3 TiI 6743.127 0.899612 1.73 +0.3 +0.1 +0.3 +0.3 +0.1 +0.2 TiII 5418.751 1.581911 2.13 +0.4 +0.0 +0.1 +0.0 +0.1 +0.0 TiII 6491.580 2.061390 2.10 +0.4 +0.2 +0.3 +0.3 +0.5 +0.1 TiII 6559.576 2.047844 2.35 +0.4 +0.3 +0.3 +0.2 +0.5 +0.3 TiII 6606.970 2.061390 2.85 +0.4 +0.2 +0.2 +0.25 +0.1 +0.1
Table B.3. Abundances of heavy elements.
Species (Å) ex(eV) gfadopted 2115 2461 2939 3514 5037 5485 HP1-2 HP1-3
EuII 6645.064 1.379816 +0.12 +0.50 +0.50 +0.70 +0.60 +0.60 +0.55 +0.40 +0.55 BaII 6141.713 0.703636 0.076 +0.50 +0.30 +0.50 0.10 +0.70 +0.50 +0.50 +0.50 BaII 6496.897 0.604321 0.32 +0.50 +0.30 +0.50 +0.10 +0.80 +0.40 +0.80 +0.45 LaII 6262.287 0.403019 1.60 +0.40 +0.30 +0.50 +0.30 +0.10 – – – LaII 6320.376 0.172903 1.56 +0.30 +0.00 +0.30 +0.25 +0.00 – 0.30 0.15 LaII 6390.477 0.321339 1.41 +0.30 +0.40 +0.50 +0.40 +0.10 +0.10 0.00 0.20 YI 6435.004 0.065760 0.82 +0.20 +0.30 +0.20 +0.50 +0.20: +0.50 0.15 – YII 6795.414 1.738160 1.19 +0.1 +0.30 +0.30 +0.20 +0.00 – 0.15 0.10 ZrI 6127.475 0.153855 1.18 +0.10 +0.00 +0.40 +0.40 +0.20 – – +0.10 ZrI 6134.585 0.00 1.43 +0.10 +0.50 +0.30 – +0.20 – +0.20: 0.00 ZrI 6140.535 0.00 +0.10 – +0.50 – – – (+0.70) – ZrI 6143.252 0.070727 1.50 +0.10 +0.10 +0.40 +0.00 +0.00 +0.30: – 0.15 SrI 6503.989 2.258995 +0.26 +0.30 – +0.30: – +0.30:: – +0.50 +0.30: SrI 6791.016 1.775266 0.73 +0.80: – – – +0.50:: – – –
